Processes in Protoplanetary Disks Phil Armitage Colorado

Processes in Protoplanetary Disks 1.Disk structure 2.Disk evolution 3.Turbulence 4.Episodic accretion 5.Single particle evolution 6.Ice lines and persistent radial structure 7.Transient structures in disks 8.Disk dispersal

The central problem The gas orbital velocity is accurately Keplerian Specific angular momentum is robustly an increasing function of radius Even though lowest energy state favors gas accreting on to the star, angular momentum conservation forbids it

The central problem Hernandez et al. ‘07 Consistent with long observed disk lifetimes – disks are quasi-equilibrium structures that evolve slowly compared to dynamical time scale

The central problem Redistribution of angular momentum within disk Loss of angular momentum in a wind (1) (2) if field line is like a rigid wire “Viscous” disk: angular momentum mixed by internal turbulence Not mutually exclusive!

Classical disks Lynden-Bell & Pringle ’74; Shakura & Sunyaev ’73 theory: disk is geometrically thin (h/r << 1), axisymmetric, planar angular momentum redistribution is modeled as a Navier-Stokes shear viscosity (kinematic viscosity ) continuity + angular momentum conservation specification of the torque G – local, scales linearly with shear

Keplerian potential specializes to… Diffusive evolution of surface density  Viscous time scale: Green’s function solution: mass flows to r = 0, while angular momentum carried by tail of mass to infinity

In steady-state, if: Also have explicit self-similar solution: Simple model to fit to observations

How applicable is classical disk theory? Angular momentum transport is not due to real “viscosity” ~km s -1 ~10 cm However, obtain same one-dimensional evolution equation if transport is due to an average turbulent stress, provided it is locally defined. e.g. for a fluid with magnetic fields, transport from fluid (“Reynolds”) and magnetic (“Maxwell”) stress Balbus & Papaloizou ‘99

How applicable is classical disk theory? Things will go wrong if we try to apply the theory when: transport mechanism is non-local (e.g. self-gravity when M disk is not much smaller than M * ) mass loss (e.g. from photoevaporation) occurs on a time scale < viscous time scale 1D situations where  far from Keplerian time scales shorter than correlation time for turbulence any 2D or 3D situation (warps, eccentric disks, meridional circulation)…

 -model disks Can make a predictive theory if we can write as a function of other disk parameters (T, r, , x e …) Shakura-Sunyaev ‘73  -prescription For  assumed constant, one parameter description of protoplanetary disk evolution Identify disk lifetime with the viscous time at outer edge t = 1 Myr at 30 AU, (h / r) = 0.05  = 0.01

Bell et al. ‘97 accretion rate M Sun yr -1 If irradiation dominates, with fixed T ~ r -1/2, then an  -disk is equivalent to ~ r (since  =  c s 2 /  ) An  model predicts the time-varying radial (and vertical) structure for any accretion rate e.g. snow line near 4 AU for this model

We can always choose to express the efficiency of angular momentum transport in terms of   -disk theory is useful if it encodes the “leading order” dependence of the stress on the local disk properties, i.e. so that  is a slowly varying function of , r etc Various caveats:  likely a strong function of T, , if transport is due to MHD processes vertical structure also depends on how accretion energy is distributed vertically… even more uncertain for comparison against observations, reducing a possibly complex function to one number

Star-disk interactions For a weakly-magnetized star: boundary layer Classical theory: point in flow at where d  /dr = 0 viscous stress vanishes disk has a boundary condition of zero torque star accretes gas with high angular momentum kinetic energy of disk is dissipated in narrow boundary layer, expected to be hot and luminous

Belyaev et al. ‘13 Boundary layer models are sensitive to the nature of disk angular momentum transport: d  / dr has opposite sign Boundary layer flow is not unstable to the magnetorotational instability, rather evidence for transport by acoustic waves (non-local, not a “viscosity” at all!)

In protostellar accretion, boundary layers occur at high accretion rate (FU Orionis objects) Kley & Lin ‘96 Radiation hydrodynamics likely important in determining the structure of the boundary layer Where is accretion energy released in FU Ori boundary layers if waves transport angular momentum? What is structure of circumplanetary disk boundary layers?

At low accretion rates, expect magnetospheric accretion Suppose vertical field at disk surface is a dipole, toroidal component similar Then magnetic torque on surface of disk Time scale for stellar torque to drive inflow is shorter than viscous time inside some magnetospheric radius r m Simulation: Dyda et al. ‘15

Very rough, but weak function due to rapid dipole fall off For kG fields, M Sun yr -1, typically r m = R Sun Consequences: gas accretes along magnetic field lines (free-fall, accretion shock on surface) magnetic field allows star to exert a non-zero torque on disk inner edge (in principle, star may spin down) innermost disk is missing

Interaction between disk and stellar field close to r m favorable location for launching jets Dyda et al.’15

Another non-zero torque case for circumbinary disks Assume gravitational torques from binary completely forbid inflow through some inner radius r = r in In 1D: Analytic Green’s function solution describing decretion disk disk L increases due to binary torque both the mass and angular momentum eventually flow outward energy / angular momentum comes from binary, which shrinks

Analytic solution Pringle ’91 Prototype for “Type II” planetary migration

Even in binaries, gas can flow across barrier and be accreted / form smaller disk around individual stars Artymowicz & Lubow ‘96 Simulation Cuadra et al. ‘09 Critical physics for massive planet growth and migration How does this operate with realistic disk turbulence?